The Model Photosphere (Chapter 9)
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Transcript of The Model Photosphere (Chapter 9)
The Model Photosphere (Chapter The Model Photosphere (Chapter 9)9)
• Basic Assumptions• Hydrostatic Equilibrium• Temperature Distributions• Physical Conditions in Stars – the
dependence of T(), Pg(), and Pe() on effective temperature and luminosity
Basic Assumptions in Stellar AtmospheresBasic Assumptions in Stellar Atmospheres
• Local Thermodynamic Equilibrium– Ionization and excitation correctly described by
the Saha and Boltzman equations, and photon distribution is black body
• Hydrostatic Equilibrium– No dynamically significant mass loss– The photosphere is not undergoing large scale
accelerations comparable to surface gravity– No pulsations or large scale flows
• Plane Parallel Atmosphere– Only one spatial coordinate (depth)– Departure from plane parallel much larger than
photon mean free path– Fine structure is negligible (but see the Sun!)
Hydrostatic EquilibriumHydrostatic Equilibrium• Consider an element of gas
with mass dm, height dx and area dA
• The upward and downward forces on the element must balance:
PdA + gdm = (P+dP)dA• If is the density at location
x, thendm= dx dA dP/dx = g
• Since g is (nearly) constant through the atmosphere, we set
g = GM/R2
P
P+dP
x+dx
x
gdm
dP/dx = g
In Optical DepthIn Optical Depth• Since d=dx
• and dP=gdx
CLASS PROBLEM:• Recall that for a gray atmosphere,
For =0.4, Teff=104, and g=GMSun/RSun
2, compute the pressure, density, and depth at =0, ½, 2/3, 1, and 2. (The density and pressure equal zero at =0 and k =1.38 x 10-16 erg K-1)
dP/d = g/
)3
2(
4
3 44 TeffT
Estimate Teff, log g, & DepthEstimate Teff, log g, & Depth
5000 T Pe Pg 5000
1.0E-5 6896 2.56E+1 4.79E+3 2.69E-1
5.0E-4 6971 1.16E+2 7.20E+4 1.06E 0
2.0E-3 7049 2.09E+2 1.75E+5 1.81E 0
1.0E-2 7179 4.27E+2 4.74E+5 3.46E 0
4.0E-2 7379 8.93E+2 1.08E+6 6.56E 0
8.0E-2 7556 1.42E+3 1.58E+6 9.62E 0
2.0E-1 7925 3.40E+3 2.48E+6 1.79E+1
5.0E-1 8601 8.90E+3 3.58E+6 4.01E+1
8.0E-1 9114 1.74E+4 4.16E+6 6.67E+1
1.60E 0 10182 5.31E+4 4.93E+6 1.58E+2
3.0E 0 11481 1.53E+5 5.49E+6 3.89E+2
6.0E 0 13228 4.47E+5 5.94E+6 1.13E+3
Wehrse Model, Teff=10000, log g=8Wehrse Model, Teff=10000, log g=85000 T Pe Pg 5000
2.00E-03 7925 5.80E+02 8.82E+04 3.74E+00
6.00E-03 8064 9.80E+02 1.71E+05 5.48E+00
1.00E-02 8129 1.24E+03 2.33E+05 7.14E+00
2.00E-02 8208 1.70E+03 3.54E+05 9.38E+00
6.00E-02 8414 3.06E+03 6.78E+05 1.55E+01
1.00E-01 8571 4.34E+03 9.01E+05 2.06E+01
2.00E-01 8920 7.70E+03 1.28E+06 3.28E+01
5.00E-01 9600 2.03E+04 1.89E+06 7.16E+01
7.00E-01 10040 3.07E+04 2.13E+06 1.01E+02
8.00E-01 10223 3.68E+04 2.22E+06 1.18E+02
1.00E+00 10544 4.98E+04 2.37E+06 1.53E+02
1.60E+00 11377 9.77E+04 2.66E+06 2.83E+02
2.00E+00 11831 1.37E+05 2.78E+06 3.95E+02
3.00E+00 12759 2.41E+05 2.96E+06 7.15E+02
4.00E+00 13476 3.50E+05 3.07E+06 1.09E+02
6.00E+00 14278 5.01E+05 3.23E+06 1.67E+02
8.00E+00 15413 7.41E+05 3.32E+06 2.71E+02
In Integral Form -In Integral Form -
• The differential form:
• x Pg½
(where 0 is at a reference wavelength, typically 5000A)
• Then integrate:
/gd
dP
00
2
1
2
1
d
gPdPP ggg
3/2
0
log
0
2/10
3
2
000
2
1
0 loglog2
3
2
3)(
00
tde
Ptg
dt
PgP gg
g
ProcedureProcedure
• Guess at Pg()
• Guess at T()
• Do the integration, computing at each level from T and Pe
• This gives a new Pg()
• Interate until the change in Pg() is small
The T(The T() Relation) Relation
• In the Sun, we can use– Limb darkening or– The variation of with wavelength
to get the T() relation• Limb darkening can be described from:
• We have already considered limb darkening in the gray case, where
cossec),0( sec
0badeSI
baS
The Solar Limb The Solar Limb DarkeningDarkening
cossec),0( sec
0badeSI
The Solar T(The Solar T() Relation) Relation
• So one can measure I(0,) and solve for S()
• Assuming LTE (and thus setting S()=B(T)) gives us the T() relation
• The profiles of strong lines also give information about T() – different parts of a line profile are formed at different depths.
The T(The T() Relation in Other Stars) Relation in Other Stars
• Use a gray atmosphere and the Eddington approximation
• More commonly, use a scaled solar model:
• Or scale from published grid models• Comparison to T(t) relations iterated
through the equation of radiative equilibrium for flux constancy suggests scaled models are close
SunSun
TTeff
TeffT )()( *
*
Comparing T(Comparing T()’s at Teff=4000, log )’s at Teff=4000, log g=2.25g=2.25
3000
3500
4000
4500
5000
5500
6000
6500
7000
7500
0.00
1
0.00
50.
02 0.1
0.4 1 2 6 10
Optical Depth
Tem
pera
ture
Scaled HM
Bell et al.
T(T() ) vs. vs. gravitgravityy
Kurucz models at 5500K
Depart at depth, similar in shallow layers
Temperature vs. MetallicityTemperature vs. Metallicity
DON’T Scale PDON’T Scale Pgg(()!)!
3000
4000
5000
6000
7000
8000
9000
1.00E+01 1.00E+02 1.00E+03 1.00E+04 1.00E+05 1.00E+06
Gas Pressure
Tem
per
atu
re (
K)
Log g = 1.0
Log g = 2.0
Log g = 3.0
Log g = 4.0
Models at 5000 K
Temperature Pressure Relation Temperature Pressure Relation with Metallicitywith Metallicity
Gas Pressure vs. MetallicityGas Pressure vs. Metallicity
Electron Pressure vs. MetallicityElectron Pressure vs. Metallicity
Computing the SpectrumComputing the Spectrum
• Now can compute T, Pg, Pe, at all (Pe=NekT)
• Does the model photosphere satisfy the energy criteria (radiative equilibrium)?
• Compute the flux from
• Express I in terms of the source function S, and adopt LTE (S =B(T))
dIF cossin2 2
0